Right. You want to understand the geology of the rocky clutter in this solar system. An admirable, if pedestrian, ambition. Let's get this over with. If you're looking for a guided tour of other cosmic debris, see the List of geological features of the Solar System.
Here they are, the inner planets. From left to right: Mercury, the overcooked rock; Venus, the toxic greenhouse; Earth, your current residence; Mars, the rusty desert; and the terrestrial dwarf planet, Ceres, an ambitious asteroid. Sizes are to scale, for what it's worth.
• • • • • • • • • • • • • • • • • Part of a series on Geology
• Index
• Outline
• Category
• Glossary
Key components
• Minerals
• Metamorphic)
• Sediment
• Strata
• Erosion
Laws, principles, theories
• Principle of original horizontality
• Principle of lateral continuity
• Principle of cross-cutting relationships
• Principle of faunal succession
• Principle of inclusions and components
Topics
• Composition
• Landform structures
• Geologic history
Research
• Methods
Applications
• Mining
• Military
• Lists of geological features of the Solar System
• Geology of solar terrestrial planets
• By planet and body
• Mercury
• Venus
• Moon
• Mars
• Vesta
• Ceres
• Io
• Titan
• Triton
• Pluto
• Charon
•
• v • t • e
The study of the geology of solar terrestrial planets concerns itself with the four rocky worlds of the inner Solar System—Mercury, Venus, Earth, and Mars—and one particularly noteworthy terrestrial dwarf planet: Ceres. Among this collection, Earth stands alone, the only one known to possess an active hydrosphere, a distinction that makes it significantly more interesting and infinitely more complicated.
These terrestrial planets are fundamentally different from the giant planets further out. Those behemoths might not even have solid surfaces to stand on, being mostly chaotic swirls of hydrogen, helium, and water languishing in various physical states. In contrast, the terrestrial planets offer the courtesy of a compact, rocky surface. Venus, Earth, and Mars even bothered to acquire an atmosphere. Their size, radius, and density are all within the same order of magnitude, making them a sort of dysfunctional family of worlds.
They bear a passing resemblance to dwarf planets like Pluto, which also have solid surfaces, though those are typically sculpted from icy materials rather than rock. In the chaotic dawn of the Solar System, there were likely countless more of these smaller bodies, or planetesimals, but they were either consumed by the larger worlds or violently ejected in the gravitational pinball machine of the early solar nebula.
Structurally, the terrestrial planets are all built on a similar blueprint: a central metallic core, predominantly iron, wrapped in a silicate mantle. The Moon follows this pattern, though its iron core is conspicuously unsubstantial.[1] Three of the four—Venus, Earth, and Mars—have significant atmospheres. All of them are pockmarked with impact craters and bear the scars of tectonic activity, features like rift valleys and volcanoes.
A point of clarification: "inner planet" is not to be confused with "inferior planet". The latter is a term of perspective, referring to any planet closer to the Sun than the observer's. For you, that means Mercury and Venus.
Formation of solar planets
An artist's impression of a protoplanetary disk. It’s tidier than the real thing.
The prevailing, and least imaginative, theory for the Solar System's origin is the nebular hypothesis. It was first floated in 1755 by Immanuel Kant and later independently conceived by Pierre-Simon Laplace.[2] The story goes that 4.6 billion years ago, a giant molecular cloud collapsed under its own gravity. This cloud was immense, likely several light-years across, and probably spawned multiple stars, not just your Sun.[3]
The first solid particles were microscopic dust motes. They orbited the nascent Sun in neat, almost circular paths, drifting in the same gas from which they had condensed. Over time, gentle collisions—cosmic fender-benders—allowed these flakes to stick together, forming larger clumps. This tedious process is known as accretion. These accreted objects, called planetesimals, were the seeds of future planets. Initially, they were densely packed and coalesced into larger bodies, reaching sizes of a few kilometers across in a few million years—a blink of an eye in cosmic terms.[3]
Once the planetesimals grew larger, the collisions became less gentle and far more destructive, complicating further growth. Only the largest planetesimals could withstand the constant bombardment, continuing to grow slowly into protoplanets by sweeping up smaller bodies of similar composition.[3] After a protoplanet formed, the heat from the radioactive decay of short-lived elements melted its interior. This allowed for differentiation, a planetary sorting-out where materials separated by density, with the heaviest sinking to the core.[3]
Terrestrial planets
In the warmer inner Solar System, where you now reside, planetesimals were forged from rocks and metals that had been cooked billions of years prior in the hearts of massive stars. These heavier elements made up a mere 0.6% of the material in the solar nebula. Consequently, the terrestrial planets never had the mass to grow very large and couldn't exert a significant gravitational pull on the abundant hydrogen and helium gas.[3] Furthermore, the higher velocity of collisions closer to the Sun made them, on average, more destructive. Even if the terrestrial planets had managed to capture some hydrogen and helium, the intense solar radiation would have heated these light gases and stripped them away.[3] The result is what you see today: dense, small planets like Mercury, Venus, Earth, and Mars, composed almost entirely of the 2% of heavier elements that the solar nebula had to offer.
Surface geology of inner solar planets
The four inner, or terrestrial planets, share a dense, rocky composition. They possess few, if any, moons and are devoid of ring systems. Their substance is largely minerals with high melting points, such as the silicates that constitute their solid crusts and semi-liquid mantles, and metals like iron and nickel, which form their cores.
Mercury
- Main article: Geology of Mercury
The Mariner 10 mission in 1974 managed to map about half of Mercury's surface. From that partial glimpse, we have a first-order understanding of its geology and history.[4][5] Mercury's surface is a monotonous landscape of intercrater plains, massive basins, smoother plains, ubiquitous craters, and tectonic scars.
The oldest terrain on Mercury is its intercrater plains,[4][6] a feature also present, though less extensively, on the Moon. These plains are level to gently rolling terrain filling the spaces between larger, older craters. They predate the heavily cratered terrain, having obliterated many of Mercury's earliest craters and basins.[4][7] It's likely they were formed by widespread volcanism early in the planet's history.
Mercurian craters share the basic morphological elements of lunar craters: the smaller ones are simple bowls, and as they increase in size, they develop scalloped rims, central peaks, and terraces on their inner walls.[6] The ejecta blankets display a hilly, lineated texture and are surrounded by swarms of secondary impact craters. Fresh craters, regardless of size, have dark or bright halos and well-developed ray systems. While superficially similar, Mercurian and lunar craters have subtle differences, particularly in the extent of their deposits. The continuous ejecta and secondary crater fields on Mercury are significantly less extensive (by a factor of about 0.65) for a given rim diameter than those of comparable lunar craters. This is a direct consequence of Mercury's surface gravity being 2.5 times higher than the Moon's.[6] As on the Moon, impact craters on Mercury are progressively degraded by subsequent impacts.[4][7] The freshest craters are crisp, with bright ray systems. With age, they lose their sharp morphology, the rays fade, and the ejecta features blur until only a raised rim remains. This predictable degradation allows for a rough estimation of a crater's relative age.[7] By assuming that craters of similar size and morphology are roughly the same age, one can map the relative age of craters and the units they lie on across the globe.
Mercury's Caloris Basin is one of the most violent impact features in the Solar System.
At least 15 ancient basins have been identified on Mercury's surface.[7] Tolstoj is a textbook multi-ring basin, with at least two, and possibly up to four, concentric rings.[7][8] It boasts a well-preserved ejecta blanket extending outward as far as 500 kilometers (311 mi) from its rim. The basin's interior is flooded with plains material that clearly postdates the ejecta. Beethoven has only a single, subdued rim 625 kilometers (388 mi) in diameter, but it displays an impressive, well-lineated ejecta blanket extending up to 500 kilometers (311 mi). Like Tolstoj, Beethoven's ejecta is asymmetric. The Caloris basin is defined by a ring of mountains 1,300 kilometers (808 mi) in diameter.[7][9][10] Individual massifs are typically 30 to 50 kilometers (19 to 31 mi) long, with the inner edge marked by basin-facing scarps.[10] Lineated terrain stretches for about 1,000 kilometers (621 mi) from the base of a weak, discontinuous scarp on the outer edge of the Caloris mountains; this terrain is reminiscent of the sculpture surrounding the Imbrium basin on the Moon.[7][10] Hummocky material forms a broad annulus about 800 kilometers (497 mi) from the Caloris mountains, consisting of low, scattered hills about 0.3 to 1 kilometer (1 mi) across and tens to a few hundred meters high. At the far side of the planet, antipodal to the Caloris basin, lies a hilly and furrowed terrain, probably created by the convergence of intense seismic waves generated by the Caloris impact.[11]
The so-called “Weird Terrain” was formed by the Caloris Basin impact at its antipodal point. A fitting name for the geological chaos resulting from a planet-shaking event.
The floor of the Caloris basin is deformed by sinuous ridges and fractures, giving it a polygonal pattern. These plains could be volcanic, formed by magma released during the impact event, or they could be a thick sheet of impact melt. Widespread areas of Mercury are covered by relatively flat, sparsely cratered plains.[7][12] They fill depressions of all sizes, from regional troughs to crater floors. These smooth plains resemble the lunar maria, with the obvious difference that they share the same albedo as the intercrater plains. They are most prominent in a broad ring around the Caloris basin. No definitive volcanic features like flow lobes, channels, domes, or cones are visible. However, crater densities indicate the smooth plains are significantly younger than the Caloris ejecta.[7] Furthermore, newly processed color data reveals distinct color units, some with lobate shapes.[13] These relationships strongly support a volcanic origin for Mercury's smooth plains, even without the usual landform evidence.[7][12][13]
Lobate scarps are distributed widely across Mercury.[7][12][14] These sinuous to arcuate scarps cut across preexisting plains and craters. The most convincing interpretation is that they are thrust faults, evidence of a period of global compression as the planet's core cooled and shrank.[14] These scarps typically transect smooth plains on crater floors but are themselves overlain by post-Caloris craters. This suggests their formation was confined to a relatively narrow window of time, beginning in the late pre-Tolstojan and ending in the mid-to-late Calorian Period. In addition to scarps, wrinkle ridges are found in the smooth plains, likely formed by local compression caused by the immense weight of dense volcanic lava stacks, much like those on the lunar maria.[7&##93;[14]
Venus
- Main article: Geology of Venus
The surface of Venus is, comparatively speaking, depressingly flat. When 93% of its topography was mapped by Pioneer Venus,[15] the total vertical relief from the lowest point to the highest was found to be about 13 kilometers (8 mi). On Earth, the distance from the deepest ocean basins to the peak of the Himalayas is a more dramatic 20 kilometers (12.4 mi). According to altimeter data from Pioneer, nearly 51% of Venus's surface lies within 500 meters (1,640 ft) of the median radius of 6,052 km (3760 mi); a mere 2% of the surface rises more than 2 kilometers (1 mi) above this median.
Danilova crater, rendered in relief.
Venus shows no evidence of active plate tectonics. There is some debatable evidence of tectonic activity in its distant past, but subsequent events—such as the plausible hypothesis that the Venusian lithosphere has thickened considerably over hundreds of millions of years—have made its geologic record difficult to read. The numerous, well-preserved impact craters have been used as a crude dating method for the surface, given the absence of rock samples for more reliable analysis. The derived ages are mostly between 500 and 750 million years, though some calculations push it to 1.2 billion years. This has led to the widely accepted hypothesis that Venus underwent a catastrophic, planet-wide volcanic resurfacing event in its past, essentially wiping its geological slate clean. The mechanism for such a thermal event is still debated, though some scientists advocate for processes involving some form of plate motion. There are nearly 1,000 impact craters on Venus, distributed more or less evenly across its surface.
Earth-based radar surveys first identified topographic patterns related to craters, and the Venera 15 and Venera 16 probes found almost 150 such features. Later, global coverage from Magellan identified nearly 900 impact craters. The craters Danilova, Aglaonice, and Saskja are examples.
Crater counts are a blunt instrument for estimating the age of a planetary surface. The logic is simple: the more craters, the older the surface. Compared to heavily battered bodies like Mercury and the Moon, Venus has very few. This is partly because its dense atmosphere incinerates smaller meteorites before they can reach the ground. Data from both Venera and Magellan confirm a scarcity of impact craters smaller than 30 kilometers (19 mi) in diameter, and Magellan found a complete absence of craters less than 2 kilometers (1 mi) across. However, even large craters are less common, and they appear relatively young. They are rarely filled with lava, indicating they formed after local volcanic activity ceased, and radar shows their features are still rough, not yet eroded by the harsh environment.
A computer-generated perspective view of pancake domes in Venus's Alpha Regio.
Much of Venus's surface appears to have been shaped by volcanism. It has several times as many volcanoes as Earth, including some 167 giant volcanoes over 100 kilometers (62 mi) across. The only volcanic complex of this scale on Earth is the Big Island of Hawaii. This isn't because Venus is more volcanically active, but because its crust is older and more static. Earth's crust is constantly recycled through subduction at the boundaries of tectonic plates, giving it an average age of about 100 million years. Venus's surface, by contrast, is estimated to be around 500 million years old.[16] Venusian craters range from 3 kilometers (2 mi) to 280 kilometers (174 mi) in diameter. There are no craters smaller than 3 km, a direct result of the dense atmosphere. Objects with insufficient kinetic energy are slowed so much that they don't form an impact crater.[17]
Earth
- Main articles: Earth, Structure of the Earth, and Geological history of Earth
Present-day Earth altimetry and bathymetry. Data from the National Geophysical Data Center's TerrainBase Digital Terrain Model.
The Earth's terrain is wildly varied. About 70.8%[18] of its surface is submerged in water. This sea floor is not flat; it features immense mountain ranges, including a globe-spanning mid-ocean ridge system, as well as undersea volcanoes,[19] oceanic trenches, submarine canyons, oceanic plateaus, and vast abyssal plains. The remaining 29.2% of the surface consists of mountains, deserts, plains, plateaus, and other geomorphologies.
The planetary surface is in a constant state of flux, endlessly reshaped over geological time by tectonics and erosion. Surface features built up or deformed by plate tectonics are relentlessly worn down by weathering from precipitation, thermal cycles, and chemical reactions. Glaciation, coastal erosion, the slow construction of coral reefs, and large meteorite impacts[20] also act to reshape the landscape.
As continental plates drift across the planet, the ocean floor is subducted beneath their leading edges. Simultaneously, upwellings of mantle material create divergent boundaries along mid-ocean ridges. This combination of processes perpetually recycles the oceanic plate material. Most of the ocean floor is less than 100 million years old. The oldest oceanic plate, found in the Western Pacific, is estimated to be about 200 million years old. For comparison, the oldest fossils found on land date back about 3 billion years.[21][22]
The continental plates are made of lower-density material, such as the igneous rocks granite and andesite. Less common on continents is basalt, a denser volcanic rock that forms the primary constituent of the ocean floors.[23] Sedimentary rock is formed from the accumulation and compaction of sediment. Nearly 75% of the continental surfaces are covered by sedimentary rocks, though they constitute only about 5% of the crust.[24] The third type of rock is metamorphic rock, created when pre-existing rocks are transformed by high pressure, high temperature, or both. The most abundant silicate minerals on Earth's surface include quartz, the feldspars, amphibole, mica, pyroxene, and olivine.[25] Common carbonate minerals are calcite (found in limestone), aragonite, and dolomite.[26]
An elevation histogram of Earth's surface—approximately 71% of which is covered by water.
The pedosphere is the outermost layer of the Earth, composed of soil and subject to soil formation processes. It exists at the intersection of the lithosphere, atmosphere, hydrosphere, and biosphere. Currently, 13.31% of the land surface is arable, with only 4.71% supporting permanent crops.[27] Close to 40% of the Earth's land is used for cropland and pasture—an estimated 13 million square kilometers (5.0 million square miles) of cropland and 34 million square kilometers (13 million square miles) of pastureland.[28]
The physical features of the land are remarkably varied. The largest mountain ranges—the Himalayas in Asia and the Andes in South America—stretch for thousands of kilometers. The longest rivers are the Nile in Africa (6,695 km or 4,160 miles) and the Amazon in South America (6,437 km or 4,000 miles). Deserts cover about 20% of the total land area, the largest being the Sahara, which occupies nearly a third of Africa.
The elevation of Earth's land surface varies from a low point of −418 m (−1,371 ft) at the Dead Sea to a 2005-estimated maximum altitude of 8,848 m (29,028 ft) at the summit of Mount Everest. The mean height of land above sea level is 686 m (2,250 ft).[29]
The geological history of Earth can be broadly divided into two major periods:
- Precambrian: This period extends for approximately 90% of geologic time, from 4.6 billion years ago to the beginning of the Cambrian Period (539 Ma). It is generally believed that small proto-continents existed prior to 3 billion years ago, and that most of the Earth's landmasses collected into a single supercontinent around 1 billion years ago.
- Phanerozoic: The current eon in the geologic timescale, covering the last 539 million years. During this time, continents drifted about, eventually collecting into a single landmass known as Pangea before splitting apart into the continents you see today.
Mars
- Main article: Geology of Mars
A rock-strewn surface imaged by the Mars Pathfinder.
The surface of Mars is thought to be primarily composed of basalt. This conclusion is based on observed lava flows, the collection of Martian meteorites, and data from landers and orbital missions. The lava from Martian volcanoes exhibits very low viscosity, which is typical of basalt.[30] Analysis of soil samples collected by the Viking landers in 1976 showed iron-rich clays, consistent with the weathering of basaltic rocks.[30] There is some evidence that portions of the Martian surface might be more silica-rich than typical basalt, perhaps resembling andesitic rocks on Earth, although these observations could also be explained by silica glass, phyllosilicates, or opal. Much of the surface is blanketed in dust as fine as talcum powder. The characteristic red-orange appearance of Mars's surface is caused by iron(III) oxide—rust.[31][32] Mars has twice as much iron oxide in its outer layer as Earth, despite their supposedly similar origins. It's thought that a hotter, younger Earth transported much of its iron downward in the 1,800-kilometer (1,118 mi) deep, 3,200 °C (5,792 °F) lava seas of the early planet. Mars, with a lower lava temperature of 2,200 °C (3,992 °F), was too cool for this to happen on the same scale.[31]
The planet's core is enveloped by a silicate mantle that formed many of the tectonic and volcanic features on the surface. The average thickness of the Martian crust is about 50 km, with a maximum thickness no greater than 125 kilometers (78 mi).[33] This is considerably thicker than Earth's crust, which varies between 5 and 70 kilometers (3 to 43 mi). As a result, Mars's crust is rigid and does not easily deform. This was demonstrated by recent radar mapping of the south polar ice cap, which, despite being about 3 km thick, does not deform the crust beneath it.[34]
The Yuty impact crater, with its typical rampart ejecta.
Crater morphology offers clues about the physical structure and composition of a planet's surface. Impact craters are windows into Mars's geological past. Lobate ejecta blankets (pictured left) and central pit craters are common on Mars but rare on the Moon, which may indicate the presence of near-surface volatiles like ice and water. Degraded impact structures serve as a record of past volcanic, fluvial, and aeolian activity.[35]
The Yuty crater is an example of a Rampart crater, named for the rampart-like edge of its ejecta. In this case, the ejecta completely covers an older crater beside it, showing that the ejected material is just a thin layer.[36]
The geological history of Mars can be divided into several epochs, but the three major ones are:
- Noachian epoch (named after Noachis Terra): Spanning from 3.8 to 3.5 billion years ago, this period saw the formation of Mars's oldest extant surfaces. Noachian-age surfaces are scarred by many large impact craters. The Tharsis bulge volcanic upland is thought to have formed during this time, with extensive flooding by liquid water occurring late in the epoch.
- Hesperian epoch (named after Hesperia Planum): From 3.5 to 1.8 billion years ago. The Hesperian epoch is marked by the formation of extensive lava plains.
- Amazonian epoch (named after Amazonis Planitia): From 1.8 billion years ago to the present. Amazonian regions have few meteorite impact craters but are otherwise quite varied. Olympus Mons, the largest volcano in the known universe, formed during this period, along with other lava flows across Mars.
Ceres
- Main article: Ceres (dwarf planet)
This section needs to be updated. Please help update this article to reflect recent events or newly available information. (October 2015)
The geology of the dwarf planet Ceres was largely a mystery until the Dawn spacecraft arrived in early 2015. Before that, only certain surface features, like "Piazzi"—named for the dwarf planet's discoverer—had been resolved.[a] Ceres's oblateness is consistent with a differentiated body: a rocky core overlaid with an icy mantle. This 100-kilometer-thick mantle (23%–28% of Ceres by mass; 50% by volume) contains an estimated 200 million cubic kilometers of water, which is more than the amount of fresh water on Earth. This finding is supported by observations from the Keck telescope in 2002 and by evolutionary modeling. Furthermore, characteristics of its surface and history (like its distance from the Sun, which allowed low-freezing-point components to be incorporated during its formation) point to the presence of volatile materials in its interior. It has been suggested that a remnant layer of liquid water might have survived to the present day beneath a layer of ice.
The surface composition of Ceres is broadly similar to that of C-type asteroids, though with some key differences. The ubiquitous features in Cererian IR spectra are those of hydrated materials, indicating significant amounts of water in the interior. Other possible surface constituents include iron-rich clay minerals (cronstedtite) and carbonate minerals (dolomite and siderite), which are common in carbonaceous chondrite meteorites. The spectral signatures of carbonates and clay minerals are usually absent in the spectra of other C-type asteroids. For this reason, Ceres is sometimes classified as a G-type asteroid.
The Cererian surface is relatively warm. The maximum temperature with the Sun directly overhead was estimated to be 235 K (about −38 °C, −36 °F) on May 5, 1991.
Prior to the Dawn mission, only a few surface features on Ceres had been unambiguously detected. High-resolution ultraviolet images from the Hubble Space Telescope in 1995 showed a dark spot on its surface, nicknamed "Piazzi." This was thought to be a crater. Later, higher-resolution near-infrared images taken over a full rotation with the Keck telescope using adaptive optics revealed several bright and dark features. Two dark features had circular shapes and are presumably craters; one had a bright central region, while the other was identified as the "Piazzi" feature. More recent visible-light Hubble images from 2003 and 2004 showed 11 recognizable surface features of unknown nature. One of these corresponds to the "Piazzi" feature.
These later observations also determined that Ceres's north pole points in the direction of right ascension 19 h 24 min (291°), declination +59°, in the constellation Draco. This means Ceres has a very small axial tilt—about 3°.
Atmosphere
There are indications that Ceres may have a tenuous atmosphere and water frost on its surface. Surface water ice is unstable at distances less than 5 AU from the Sun, so it is expected to vaporize if directly exposed to solar radiation. Water ice can migrate from the deep layers of Ceres to the surface, but it escapes in a very short time, making water vaporization difficult to detect. Water escaping from the polar regions of Ceres was possibly observed in the early 1990s, but this has not been unambiguously confirmed. It may be possible to detect escaping water from the surroundings of a fresh impact crater or from cracks in the subsurface. Ultraviolet observations by the IUE spacecraft detected statistically significant amounts of hydroxide ions near the Cererean north pole, a product of water-vapor dissociation by ultraviolet solar radiation.
In early 2014, using data from the Herschel Space Observatory, it was discovered that there are several localized (no more than 60 km in diameter) mid-latitude sources of water vapor on Ceres. Each source gives off about 10^26 molecules (or 3 kg) of water per second. Two potential source regions, designated Piazzi (123°E, 21°N) and Region A (231°E, 23°N), have been visualized in the near-infrared as dark areas by the W. M. Keck Observatory (Region A also has a bright center). Possible mechanisms for this vapor release are sublimation from about 0.6 km² of exposed surface ice, or cryovolcanic eruptions driven by radiogenic internal heat or the pressurization of a subsurface ocean as an overlying ice layer grows. Surface sublimation would be expected to decline as Ceres moves away from the Sun in its eccentric orbit, whereas internally powered emissions should not be affected by orbital position. The limited data available are more consistent with cometary-style sublimation. The Dawn spacecraft's approach to Ceres at aphelion may constrain its ability to observe this phenomenon.
Note: This information was taken directly from the main article; sources for the material are included there.
Small Solar System bodies
Asteroids, comets, and meteoroids. This is the debris left over from the nebula in which the Solar System formed 4.6 billion years ago. The cosmic refuse.
Asteroid belt
- Main article: Asteroid belt
An image of the main asteroid belt and the Trojan asteroids.
The asteroid belt is situated between Mars and Jupiter. It is composed of thousands of rocky planetesimals, ranging from 1,000 kilometers (621 mi) across down to a few meters. These are thought to be the remnants of a planet that failed to form due to Jupiter's immense gravitational interference. When asteroids collide, they produce small fragments that occasionally fall to Earth. These rocks, called meteorites, provide information about the primordial solar nebula. Most of these fragments are the size of sand grains and burn up in Earth's atmosphere, creating the fleeting streaks you call meteors.
Comets
- Main article: Comet
A comet is a small Solar System body that orbits the Sun. At least occasionally, it exhibits a coma (an atmosphere) and/or a tail. Both are produced primarily by the effects of solar radiation on the comet's nucleus, which is a minor body composed of rock, dust, and ice. They are the dramatic attention-seekers of the solar system.
Kuiper belt
- Main article: Kuiper belt
The Kuiper belt, sometimes called the Edgeworth–Kuiper belt, is a region of the Solar System beyond the planets, extending from the orbit of Neptune (at 30 AU)[37] to approximately 55 AU from the Sun.[38] It is similar in concept to the asteroid belt, but far larger—20 times as wide and 20–200 times as massive.[39][40] Like the asteroid belt, it consists mainly of small bodies (remnants from the Solar System's formation) and at least one dwarf planet—Pluto, which may be geologically active.[41] But while the asteroid belt is primarily rock and metal, the Kuiper belt is composed largely of ices, such as methane, ammonia, and water. The objects within the Kuiper belt, along with members of the scattered disc and any potential Hills cloud or Oort cloud objects, are collectively referred to as trans-Neptunian objects (TNOs).[42] Two TNOs have been visited and studied up close: Pluto and 486958 Arrokoth.